chapter 8 high-altitude production of titan’s...

201

Abstract Measurements with the Cassini Ion and Neutral Mass Spectrometer (INMS) and two Cassini Plasma Spectrometer (CAPS) sensors, the Ion beam Spectrometer (IBS) and the Electron Spectrometer (ELS), have revealed the presence of a signifi cant population of heavy hydrocarbon and nitrile species well above the homopause, with masses as large as several thousand Daltons (Da). The INMS ion and neutral spectra cover the mass range 1–100 Da. The IBS has measured positive ions up to 350 Da, while the ELS has detected concen-trations of negative ions as high as 20% of the total negatively charged ionosphere component extending to over 13,000 Da. These measurements have motivated the development of new atmospheric models and have signifi cant implications for our knowledge and understanding of Titan’s haze layers.

The existence of a thick haze obscuring Titan’s surface was inferred from remote-sensing observations at infrared and ultraviolet wavelengths during the mid-1970s (Danielson et al. 1973; Veverka 1973; Zellner 1973; Trafton 1975) and confi rmed by Voyager 1 and 2 imaging, which revealed the existence of two principal haze layers, a main layer and a thin detached layer ∼100 km above it, both merging at high northern latitudes (Smith et al. 1981, 1982). It was recog-nized early on (e.g., Danielson et al. 1973) that photochem-istry occurring in the upper atmosphere of Titan was the likely source of the haze-forming aerosols, and in the years

leading up to the Voyager encounters several laboratory experiments were performed in an attempt to synthesize materials whose properties were similar to those of the postu-lated hazes (see reviews by Chang et al. 1979 and Cabane and Chassefi ère 1995). Substances investigated as possible candidates for the haze-forming aerosols included polymers of acetylene, ethylene, and HCN (Scattergood and Owen 1977; Podolak and Bar-Nun 1979) and “tholins,” complex organic solids, brownish in color, produced in a simulated reducing planetary atmosphere through UV irradiation and electric discharge (Khare and Sagan 1973; Sagan and Khare 1979).

Prior to the Voyager encounters, the only species known with certainty to be present in Titan’s atmosphere were CH

4

and C2H

6, although there was evidence for the presence of

C2H

2 and C

2H

4 as well (Gillett 1975). The presence of N

2,

predicted by Hunten (1977) and Atreya et al. (1978), had not yet been established, although Titan’s reddish-brown albedo suggested that nitrogen-bearing species (and/or sulfur-bearing ones) should be present in the haze aerosols (Scattergood and Owen 1977; Chang et al. 1979). The Voyagers revealed that Titan’s atmosphere consists predominantly (>90%) of molecular nitrogen (Broadfoot et al. 1981; Tyler et al. 1981) with methane as the next most abundant species and provided positive identifi cations of several hydrocarbons including C

2H

2, C

2H

4, and C

3H

8 as well as of the nitriles HCN, HC

3N,

and C2N

2 (Hanel et al. 1981, 1982; Kunde et al. 1981;

Maguire et al. 1981).During the interval between the Voyager encounters and

the arrival of Cassini in the Saturn system, several photo-chemical models were developed to describe the production of hydrocarbons and nitriles resulting from the dissociation of N

2 and CH

4 in Titan’s upper atmosphere by electron

impact (N2) and UV irradiation (CH

4) (e.g., Yung et al. 1984;

Toublanc et al. 1995; Wilson and Atreya 2004). More a number of laboratory, modeling, and theoretical studies were under-taken to investigate the formation of the haze layers and the physical, optical, and chemical properties of the aerosols in light of both the Voyager data and new remote-sensing

Chapter 8High-Altitude Production of Titan’s Aerosols

J.H. Waite, Jr., D.T. Young, J.H. Westlake, J.I. Lunine, C.P. McKay, and W.S. Lewis

J.H. Waite, Jr.,(�), D.T. Young, J.H. Westlake, and W.S. LewisSouthwest Research Institute, P.O. Drawer 28510 San Antonio, TX 78228, USAe-mail: [email protected]

J.H. Waite, Jr., J.H. WestlakeUniversity of Texas at San Antonio, One UTSA Blvd. San Antonio, TX 78249, USA

J.I. LunineLunar and Planetary Laboratory, University of Arizona, Tucson, AZ 85721, USA

C.P. MckayNASA Ames Research Center, Mail Stop 245–3, Moffett Field, CA 94035, USA

R.H. Brown et al. (eds.), Titan from Cassini-Huygens, © Springer Science +Business Media B.V. 2009

202 J.H. Waite et al.

observations (see reviews by Cabane and Chassefi ère 1995 and McKay et al. 2001). Post-Voyager experiments to synthesize aerosol analogs in the laboratory involved both the production of tholins in a simulated Titan N

2-CH

4 atmo-

sphere (e.g., Thompson et al. 1994; Coll et al. 1999) and the creation of the photopolymers of C

2H

2, C

2H

4, and HCN (Bar-

Nun et al. 1988; Scattergood et al. 1992) as well as of HC3N

and HC3N/C

2H

2 (Clarke and Ferris 1997). The spectral and

optical properties of tholins were found to be consistent with Titan’s albedo and with the refractive properties of Titan’s haze particles, suggesting that tholins are good analogs for Titan’s aerosols (Khare et al. 1984).

The ultimate sources of Titan’s aerosols are the gas-phase dissociation products of CH

4 and N

2. However, as noted by

Lebonnois et al. (2002), the transition from gas-phase compounds to solid-phase aerosols is poorly understood. They suggested three possible chemical pathways that could polymerize simple molecules to macromolecules, which are the presumed precursors to aerosol particles, producing: (1) polymers of acetylene and cyanoacetylene, (2) polycyclic aromatics, and (3) polymers of HCN and other nitriles, and polyynes. Their model suggested a total production rate of 4 × 10−14 g cm−2 s−1 and a C/N ratio of 4, in a production zone slightly lower than 200 km altitude. Wilson and Atreya (2003) considered similar pathways and concluded that the growth of polycyclic aromatic hydrocarbons (PAH) throughout the lower stratosphere could play an important role in haze formation. They suggested that the peak chemical produc-tion of haze would lie near 220 km, with a column integrated production rate of 3.2 × 10−14 g cm−2 s−1. Wilson and Atreya (2003) pointed out that the discovery of benzene in Titan’s atmosphere by ISO (Coustenis et al. 2003) favored the PAH pathway. Trainer et al. (2004) found that for particles produced from a mixture of 10% CH

4 in N

2 the results were

consistent with a large fraction of aromatics, including specifi c mass spectral peaks likely due to PAHs. However, at lower concentrations of CH

4 (1% and lower), the mass

fraction of PAHs greatly diminished, and an aliphatic pathway dominated.

Laboratory simulations also indicate a possible key role for PAHs. Khare et al. (2002) reported on an analysis of the time-dependent chemical evolution of gas phase products in a Titan simulation. They found an early dominance of aro-matic ring structures that led in the later stages of the experi-ment to the appearance of nitrile and amine compounds. Thompson et al. (1991) reported the yields of gaseous hydro-carbons and nitriles produced at pressures (1,700 Pa and 24 Pa) in a continuous-fl ow, low-dose, cold plasma discharge excited in an atmosphere consisting of 10% CH

4 and 90% N

2

at 295 K. At 1,700 Pa, 59 gaseous species including 27 nitriles were detected while at 24 Pa, 19 species are detected, including six nitriles and three other unidentifi ed N-bearing compounds. The types of molecules formed changed even

more markedly, with high degrees of multiple bonding at 24 Pa prevailing over more H-saturated molecules at 1,700 Pa. Imanaka et al. (2004) conducted a series of experiments from high (2,300 Pa) to low (13 Pa) pressure. They found an increase in the aromatic compounds and a decrease in C/N ratio in tholins formed at low pressures, indicating the presence of the nitrogen-containing polycyclic aromatic compounds in tholins formed at low pressures. They concluded that the haze layers at various altitudes might have different chemical and optical properties, but most importantly they found that there is a fundamental change in the nature of haze production between pressures above and below roughly 100 Pa.

8.1 Cassini Observations of Heavy Hydrocarbons in Titan’s Upper Atmosphere

Earlier models of the photochemistry responsible for initiating the production of complex acetylene polymers and polyaro-matic hydrocarbons (PAHs) suggested that the formation of heavy hydrocarbons such as benzene occurs primarily in the well-mixed portion of Titan’s atmosphere below the homo-pause (pressure ∼2 × 10−3 Pa near 750 km) (e.g., Wilson and Atreya 2003, 2004). Ion-neutral reactions near the ionospheric peak (pressure ∼5 × 10−6 Pa at 1,100 km) were thought to be an additional, although much weaker source of complex hydrocarbons. Thus, prior to the arrival of Cassini, it was expected that there would be little benzene or other complex hydrocarbons, and thus they would be only marginally detectable at altitudes above ∼950 km, the region sampled by the Cassini orbiter during its passes through Titan’s upper atmosphere. However, measurements with the Cassini Ion and Neutral Mass Spectrometer (INMS) and two Cassini Plasma Spectrometer (CAPS) sensors, the Ion beam Spectrometer (IBS) and the Electron Spectrometer (ELS), have revealed the presence of a signifi cant population of heavy hydrocarbon and nitrile species well above the homo-pause, with masses as large as several thousand Daltons (Da). The INMS ion and neutral spectra cover the mass range 1–100 Da (Fig. 8.1; Tables 8.1 and 8.2) (Waite et al. 2005; Magee et al. 2009). The IBS has measured positive ions up to 350 Da (Fig. 8.2; Crary et al. 2009), while the ELS has detected a concentration of negative ions as high as 20% of the total negatively charged ionospheric component extending to over 13,000 Da (Fig. 8.3; Coates et al. 2009).

The INMS is a true (quadrupole) mass spectrometer designed to measure the abundance of ion and neutral species in Titan’s upper atmosphere (Waite et al. 2004). The IBS and ELS sensors, on the other hand, measure ion fl ux as a funatic of ion energy/charge from which pseudo-mass spectra can be derived. IBS, designed for the supersonic solar wind and the

2038 High-Altitude Production of Titan’s Aerosols

Fig. 8.1 Composite mass spectra for neutrals (top) and ions (bottom) based on Cassini INMS data (black line) acquired during 17 fl ybys of Titan. Data were taken between 1,000 and 1,100 km. The mass deconvolution used to produce Table 8.1 is indicated above the spectrum and the totals are marked on the spectrum (top panel) with red dots, the top panel (from Waite et al. (2007) and is reprinted with permission from AAAS)

cold ionosphere of Titan, has relatively high energy resolu-tion (DE/E = 1.7%). ELS, which is designed to measure hot plasma electrons, has lower resolution (DE/E = 17%) but is in addition sensitive to negative ions (Waite et al. 2007; Coates et al. 2007, 2009).

As a consequence of the low temperatures of ions in Titan’s ionosphere (120–250 K between ∼950 and ∼1,600 km respectively), ion thermal velocities are small (100’s of m/s, depending on ion mass). During encounters with Titan Cassini travels at supersonic velocities (∼6 km/s) relative to the cold ionosphere, which allows the use of IBS and ELS energy/charge spectra to infer ion mass/charge regardless of

charge state or polarity. Crary et al. (2009) used INMS data combined with IBS to extract mass spectra and estimate ion temperatures and fl ow speeds. Similarly Coates et al. (2007, 2009) and Waite et al. (2007) analyzed ELS data to produce negative ion mass spectra.

The peak fl ux measured by IBS or ELS can be identifi ed with ion mass by the relationship E

i = m

i V

s/c2/2 + 8 kT

where Vs/c

is spacecraft velocity and mi is the mass of the

ith species (Crary et al. 2009). Taking T = 150 K as a char-acteristic temperature, E

i = 0.188 m

i + 0.013 eV. Carbon is

the smallest mass of interest here, so mi ≥12 Da and m

iV

s/

c2/2 >>; 8 kT. Thus E

i = m

iV

s/c2/2 to a very good approximation


Table 8.1 Neutral species mixing ratios measured by INMS in the altitude region between 1100 and 1000 km. Values are globally averaged over 20 fl ybys as reported by Magee et al. (2009)

Major speciesMixing ratio

N2

0.963 ± 0.44 × 10−314N15N 1.08 × 10–2 ± 0.06 × 10−2

CH4

2.17 × 10–2 ± 0.44 × 10−213CH

42.52 × 10–4 ± 0.46 × 10−4

H2

3.38 × 10–3 ± 0.23 × 10−3

Minor speciesMixing ratios

C2H

23.42–3.43 × 10−4

C2H

43.91–3.97 × 10−4

C2H

66.05–4.57 × 10−5

HCN 2.40–2.44 × 10−440Ar 2.14–2.26 × 10−5

CH3CCH 0.92–1.13 × 10−5

C3H

62.33–3.45 × 10−6

C3H

82.87–4.38 × 10−6

C4H

25.55–5.65 × 10−6

C2N

22.14–2.20 × 10−6

C6H

62.48–2.50 × 10−6

C2HCN 1.48–1.54× 10−6

C2H

3CN 3.46–4.39 × 10−7

C2H

5CN 1.54–2.87 × 10−7

C7H

82.51–5.37 × 10−8

and mass (in Daltons) can be inferred using mi = 5.32E

i.

Under these circumstances the energy resolution of IBS is equivalent to an effective mass resolution of 30 at 28 Da. Peaks in the IBS energy spectra can be resolved up to ∼200 Da (Fig. 8.2), while the maximum mass observed so far is ∼350 Da. By assuming maxwellian velocity distribu-tions and comparing IBS with INMS, Crary et al. (2009)

were able to correct IBS data for spacecraft potential and ion winds along the spacecraft track to obtain pseudo-mass spectra (Fig. 8.4).

Figure 8.5 shows an example of ion density calculated by Crary et al. using this technique. The agreement between the two instruments below ∼1,600 km is very good as is agreement with the Cassini Langmuir Probe measurements of total plasma density (Wahlund et al. 2009). The lack of agreement above this altitude is caused by ion heating, which invalidates the cold ion assumption used to interpret IBS data. Heavy ions >100 Da become a major constituent below ∼1,200 km and tend to increase deeper in the ionosphere, becoming as much as ∼50% of the total at 950 km, the lowest altitudes visited by Cassini (Crary et al. 2009). Altitude profi les strongly suggest that the abundance of ions >100 Da continues to increase rapidly with depth in the atmosphere (Fig. 5; also Wahlund et al. 2009).

Detailed knowledge of the abundance of heavy ions is important for understanding the chemistry of ion formation. Using data from all 14 Titan encounters studied so far, Crary et al. were able to obtain 130 mass spectra. (The number of spectra is limited by the need to swing the CAPS sensors across the Cassini ram direction every ∼60 s as is evidenced by the spacing of peaks in Fig. 8.6, which shows the percent occurrence calculated for each mass bin from 100 to 200 Da. On more recent passes such as T55, motion of the sensors has been halted to yield a much larger number of spectra in the ram direction. These have not yet been analyzed.) When taken together and examined statistically the spectra show a consistent mass peak spacing of 12–14 Da as expected for compounds consisting of carbon and nitrogen. Analysis of the abundant groups >100 Da, together with INMS data 70% of the spectra. In addition they conclude that aromatic hydrocarbons, particu-larly PAHs, are the most likely component of the positive ion spectra. Sittler et al. (2009) have advanced arguments to the effect that the negative ions observed by ELs could be fuller-enes (C60), a suggestion based on analogies to observations of fullerenes formed under laboratory conditions similar to those in Titan's atmosphere. Since neutrals and ions are cold, energy would come from hot electrons (0.1 ~ few eV) that are abundant in the ionosphere (Coates et al., 2009). Sittler et al. go on to suggest that because there are a number of sources of oxygen present in the atmosphere, particularly ~ke V oxygen and water group ions driven in to the atmosphere by corotation, oxygen might in fact be trapped in the fuller-enes, a phenomenon also observed in the laboratory. Once formed the fullerenes would settle into haze layers as dis-cussed above, eventually reaching the surface where they would represent a source of oxygen that might contribute to

Table 8.2 Ion densities measured by INMS for fi ve passes

T16 T17 T18 T21 T23

CH5+ 1.7 19.7 4.85 0.27 12.97

(0.1) (0.3) (0.18) (0.03) (0.25)C

2H

5+ 1.3 95.8 11.06 0.55 61.16

(0.2) (0.6) (0.21) (0.04) (0.87)HCNH+ 18.5 499.5 49.4 2.47 242.8

(0.3) (1.3) (0.48) (0.10) (2.74)C

3H

3+ 30.1 69 44.19 16.45 159.9

(0.4) (1.4) (1.10) (0.87) (4.48)C

4H

3+ 3.1 9.5 3.33 0.48 21.66

(0.3) (0.5) (0.27) (0.12) (1.13)C

4H

5+ 3.8 6.3 4.06 0.75 26.98

(0.3) (0.4) (0.30) (0.15) (1.38)C

6H

5+ 3.8 0.4 2.67 1.68 20.67

(0.4) (0.1) (0.23) (0.28) (1.47)C

6H

7+ 3.4 0.6 3.72 0.90 33.51

(0.4) (0.1) (0.27) (0.20) (2.26)C

7H

7+ 9.2 1.2 11.35 5.88 83.97

(0.9) (0.2) (0.50) (0.72) (5.79)Altitude (km) 950 1000 960 1000 1000LST 17.4 10.5 4.8 20.4 14.1


pre-biological chemistry. Although somewhat speculative in nature (there are many steps involved, not all of which are well established) the possibility of fullerene formation is certainly worth further investigation.

Figure 8.8 shows ELS energy-time spectrograms for four encounters. The sharp spikes in all four panels indicate the presence of cold negative ions that are seen only when ELS sweeps through the ram direction. Although ELS has much lower energy resolution than IBS, the same principles of analysis can be applied, enabling energy spectra to be con-verted into pseudo-mass spectra (Fig. 8.3). Mass resolution is limited to ∼5 at ∼16 Da, and, as for IBS, drops at higher masses. Because of the low instrument resolution and high ion masses observed by ELS, corrections for spacecraft potential and winds can be neglected.

The fi nding of abundant heavy negative ions using ELS is one of the truly surprising “discoveries” made with Cassini. Although photoelectron peaks are seen during daylight encounters, it is very clear that the peaks in the spectra iden-tifi ed as negative ions are far too narrow and unidirectional to be misidentifi ed as electrons, which are both isotropic and hot (>>1 eV, equivalent to >> 10,000 K). As Fig. 8.3 shows, the mass of negative ions extends from 17 Da to >10,000 Da. Negative ions are a permanent feature of the ionosphere, having been observed on all 23 encounters thus far during which spacecraft pointing was favorable for observations. Two very clear negative ion peaks can typically be identified in the spectra at 22 ± 4 Da and 44 ± 8 Da with a possible third peak at 82 ± 14 Da (Figs. 8 and 2 of Vuitton et al. 2007).

Fig. 8.2 IBS spectrum from fl yby T26. The match to the Cassini INMS ion spectrum below 100 Da is marked in red. Note the signifi cant ion densities above 100 Da

Fig. 8.3 Inferred mass/charge spectrum of negatively charged ions using ELS energy/charge data (energy scale is shown at top of fi gure). Spectra were taken at 953 km during the T16 pass. Upper trace in top panel shows count rate spectrum corrected for photoelectron contribution. Lower panel shows spectrum converted to differential number density (from Coates et al. (2007) )


Fig. 8.5 Comparison of INMS and IBS ion densities during the T26 encounter (Crary et al. 2009). INMS total ion density is shown in black. Using data from the IBS, Crary et al. calculated the total ion density (red), density of ions below 100 Da (blue) and density of ions heavier than 100 Da (green) (reprinted with permission from Elsevier)

Fig. 8.4 Mass spectra from the INMS (upper panel) and the IBS (lower panel) from 1,025 km during the ingress leg of the T26 encounter. The lower panel shows a best fi t to the IBS data below 100 Da using the method described in the text. Note that, although poorly resolved, mass peaks are still visible above 100 Da. The bottom panel (from Crary et al. (2009) with permission from Elsevier)

Fig. 8.6 The percent occurrence calculated for each mass bin from 100 to 200 Da (Crary et al. 2009). Total percent occurrence for each group is shown in the box above each peak (reprinted with permission from Elsevier)


Fig. 8.7 Color-coded fi gure showing the likelihood of chemical groups to be present in the high mass ion population. Probabilities are determined from the percent occurrence spectrum shown in Fig. 8.6 (reprinted from Crary et al. (2009) with permission from Elsevier)

Fig. 8.8 Energy-time spectrograms of ELS data taken during T16 to T19 and centered on closest approach to Titan (Coates et al. 2007). The prominent peaks in each panel are due to cold negative ions rammed into the instru-ment. Intense fl uxes below ∼10 eV correspond to photoelectrons from the spacecraft which disappear when spacecraft potential becomes negative in the ionosphere. Fluxes between ∼10 and 30 eV correspond to ionospheric photoelectrons


CAPS observations combining IBS and ELS data show that heavy positive and negative ions are present on every pass during the primary mission, 14 in all, where pointing was appropriate. (cf. Table 1 of Crary et al. 2009 and Table 1 of Coates et al. 2009; Wahlund et al. 2009. Analysis has not yet been completed for passes in the extended mission, which began in summer 2008.) Heavy positive ions become a signifi -cant component of the positive ion component of the iono-sphere below 1,200 km, while at the same altitude heavy negative ions begin to become a prominent fraction (∼20%) of the total negatively charged ionospheric component, presum-ably electrons (Coates et al. 2007, 2009; Waite et al. 2007).

Total negative ion density increases with decreasing alti-tude as does the maximum negative ion mass (Fig. 8.9). The latter is strongly dependent on altitude, varying from approxi-mately few 100 Da at ∼1,400 km to as much as 13,800 Da at 950 km (Coates et al. 2007, 2009). At that altitude negative ions can reach densities up ∼200 cm−3 and make up as much as 20% of all the negative charge present (Wahlund et al. 2009).

Although a particle with mass of 13,800 Da is a very large molecule indeed, it is doubtful that these are true molecules. Rather they are most likely aerosols formed by the clumping of smaller molecules. Since only mass/charge can be inferred from ELS measurements, if ion charge is >1 electron, then ion mass and size would be proportionately larger. Although simple to estimate in principle, calculating particle size depends on an assumption of density. Solid particles might have a characteristic density ∼1 g/cm3 whereas fractal parti-cles might conceivably have densities as low as 0.001 g/cm3. The estimated radius for maximum observed mass (13,800 Da) is then in the range 3.8 to 38 nm, the size of small aerosols which, it has been suggested, could be precur-

sors of larger aerosols seen at lower altitudes (see Chapter 12 and references therein).

Since size and electrical charge are important microphys-ical properties of the aerosols, estimations of size are critical to developing theories of formation and growth (Tomasko et al. 2008; Lavvas et al. 2009). At ∼1,000 km altitude the size of heavy negative ions is much smaller than the local Debye length (approximately few centimeters) leading to an esti-mated particle potential j ∼ −2.5 kT

e/e (Goertz 1989). Since

only mass/charge is known, in order to estimate particle size we need to establish ion charge. The electron energy spectrum measured by ELS is variable, at times in sunlight dominated by a photoelectron peak at a few electronvolts (Vuitton et al. 2009). In shadow, the characteristic equilib-rium electron temperature is ∼1,000 K. On the other hand, the assumption of thermal equilibrium between electrons and ions would result in a temperature of ∼150 K. These values lead to a wide range of particle potentials ranging from ∼−5 to ∼−0.01 V depending on conditions and location.

With an estimate of potential, particle charge can be esti-mated from Q = 4pe

0aj exp(−a/l

D) where a is particle radius

which we earlier estimated to be in the range of 3.8 to 38 nm depending on density. Since a


NH2− (0.77 eV) and CN− (3.86 eV) makes them candidates for

the components of the fi rst peak at ∼22 Da. (Although ion masses are 16 and 26 respectively both could be present in varying amounts.) Similar considerations suggest that the second peak might be NCN−, HNCN− or C

3H− (see Coates

et al. 2007) or C4H− (Vuitton et al. 2009). Vuitton et al. also

suggest C5N− as a possible candidate for third peak. Candidates

for heavier ions involve polyynes, nitriles, PAHs, Fullerenes (Sittler et al., 2009) and cyano-nitriles. There is no shortage of possibilities, but at present there is not enough information to narrow the selection. However details of composition mat-ter less than the fact that there is a rich soup of heavy organic compounds conducive to forming aerosols.

8.2 New Chemical Models Based on the Cassini Results

The measurements made by Cassini have motivated a new round of modeling and analysis, which has contributed to our present level of understanding of organic formation in Titan’s atmosphere. Waite et al. (2007) were the fi rst to high-light the correspondence of the ion and neutral mass spectra in the upper atmosphere and suggested that this demon-strated a strong degree of coupling in the ion-neutral chem-istry that resulted in the observed complexity of the organic compounds. Furthermore, they showed that ion neutral chemistry involving C

4H

3+ was the likely formation path of

the observed benzene and that benzene was in chemical equilibrium with its protonated ion counterpart, C

6H

7+.

Vuitton et al. (2006) had earlier postulated that most of the unexpected ion peaks measured by INMS in the T5 fl yby were the result of protonated nitrile compounds. They used this basic premise to develop a complex zero-dimensional model of the ion neutral chemistry and to infer much of the trace nitrile composition in the neutral atmosphere. Carrasco et al. (2008) examined this same ion neutral data set using two chemical reaction pruning methods coupled with a Bayesian statistical analysis of the reaction rate uncertainty to conclude that only 35 key ion molecule reactions were needed to describe the basic ion-neutral coupling and that they were not dominated by proton transfer as Vuitton et al. had implied. Rather they consisted of 32 growth reactions leading to chemical complexity (see Table 8.3) including 22 condensation (bond rearrangement reactions), 5 protona-tions, and 5 charge transfer reactions.

Using a method that accounts for diurnal variations in the energy input into the upper atmosphere, De La Haye et al. (2008) have developed a coupled ion-neutral composition model for Titan’s atmosphere over the altitude range from 600 to 2,000 km. The model demonstrates the important role played by ion-neutral chemistry in the formation of both hydrocarbons

and nitriles, although the relative importance of ion-neutral vs neutral chemistry varies with local time, altitude, and species.

More comprehensive one-dimensional models of the atmosphere that extend from the surface to the exobase and include both gases and particulate chemistry have been recently presented by Lavvas et al. (2008a,b) and Krasnopolsky (2009). Krasnopolsky assumed the temperature structure measured by HASI, whereas Lavvas et al. calculated the thermal structure self consistently, which in fact agrees with the HASI derived profi le. Both models generate the haze structure from the gaseous species photochemistry and the authors compare their results to the Cassini INMS and CIRS composition measurements in the thermosphere and stratosphere, respec-

Table 8.3 Ion neutral reactions deemed important for the upper atmosphere as assessed by Carrasco et al. (2008)

Reaction

Global rate constant, k × 10–9 cm−3 s−1

Branching ratio (br) Dk/k

CH+ + CH4 → C

2H

3+ + H

21.30 0.84 0.20

CH+2 + CH

4 → C

2H

4+ + H

21.30 0.7 0.15

CH+2 + HCN → C

2H

2N+ + H 1.80 1 0.20

CH+3 + CH

4 → C

2H

5+ 1.10 1 0.20

CH +4 + CH

4 → CH

5+ + CH

31.14 1 0.15

CH+5 + C

2H

2 → C

2H

3+ + CH

41.48 1 0.20

CH+5 + C

2H

4 → C

2H

5+ + CH

41.50 1 0.20

CH +5 + C

2H

6 → C

2H

7+ + CH

41.35 0.85 0.15

C2H +

2 + CH

4 → C

3H

5+ + H 0.89 0.79 0.10

C2H+

3 + CH

4 → C

3H

5+ + H

20.19 1 0.20

C2H+

4 + C

2H

2 → c-C

3H

3+ + CH

30.84 0.77 0.10

C2H +

5 + C

2H

2 → c-C

3H

3+ + CH

40.19 0.36 0.10

C2H+

5 + C

2H

4 → C

3H

5+ + CH

40.35 1 0.15

C2H +

5 + C

3H

8 → C

3H

7 + + C

2H

60.63 1 0.15

C2H+

5 + H

2O → H

3O+ + C

2H

41.86 1 0.65

C2H+

5 + HCN → HCNH+ + C

2H

42.70 1 0.20

C3H+ + CH

4 → C

2H

3 + + C

2H

20.87 0.9 0.20

C3H+

5 + C

2H

2 → C

5H

5+ + H

20.38 1 0.15

N(3P)+ + CH4 → CH

3+ + NH 1.15 0.53 0.35

N(3P)+ + CH4 → CH

4+ + N 1.15 0.05 0.35

N(3P)+ + CH4 → HCNH+ + H

21.15 0.32 0.35

N(3P)+ + C2H

2 → C

2H

2+ + N 1.50 0.70 0.15

N(3P)+ + C2H

2 → CNC+ + H

21.50 0.15 0.15

N(3P)+ + C2H

4 → C

2H

2N+ + H

21.45 0.05 0.24

N+2 + H

2 → N

2H+ + H 1.70 1 0.24

N+2 + CH

4 → CH

2+ + H

2 + N

21.14 0.08 0.15

N+2 + CH

4 → CH

3+ + H + N

21.14 0.89 0.15

N+2 + C

2H

2 → C

2H

2+ + N

20.46 1 0.25

N+2 + C

2H

4 → C

2H

3+ + H + N

21.30 0.63 0.15

N2H+ + CH

4 → CH

5+ + N

20.89 1 0.30

H3O+ + HCN → HCNH+ + H

2O 3.80 1 0.15

HCNH+ + CH3CN →

C2H

4N+ + HCN

3.80 1 0.20

HCNH+ + NH3 → NH

4+ + HCN 2.30 1 0.30

CNC+ + C2H

2 → C

3H+ + HCN 0.80 0.91 0.40

CNC+ + C2H

4 → C

2H

2N+ + C

2H

21.00 0.25 0.20

Source: Adapted from Carrasco et al. (2008), with permission from Wiley Interscience. © Wiley Interscience, New York, 2008.


tively, and to the DISR measurements of the haze from the Huygens probe. Lavvas et al. identifi ed a new source of haze particles from 500 to 900 km based on nitrile-hydrocarbon copolymer chemistry, but did not include ion-neutral chemistry and thus failed to identify the source of the high latitude mac-romolecules observed by CAPS and INMS. Haze formation in their model is dominated by nitrile and aromatic chemistry below 300 km. The total rate of precipitation they calculate is 1.27 × 10−14 g cm−2 s−1 (4 kg cm−2 Gy−1), which is in the range estimated by McKay et al. (2001), 0.5–2 × 10−14 g cm−2 s−1 (1.6–6.3 kg cm−2 Gy−1). Krasnopolsky (2009) included ion-neutral chemistry, ambipolar diffusion, and atmospheric escape as well as both positive and negative ion chemistry that to fi rst order agrees with the CAPS and INMS results. The major haze production in this case is via the reactions C

6H + C

4H

2 (polyyne chemistry), C

3N + C

4H

2 (copolymer

chemistry), and the condensation of hydrocarbons below 100 km fi rst suggested by Hunten (2006). Overall, the esti-mated precipitation rate is equal to 1.2–2.2 × 10−14 g cm−2 s−1 (4–7 kg cm−2 Gy−1).

Tables 8.4 and 8.5 compare the models discussed above to INMS measurements at 1,050 km and to CIRS measure-ments at 300 km, respectively. The modeling community before 2008 in general over-predicted the abundance of the C2 group, possibly due to overproduction of C

2H

5+ and

C3H

5+ in the case of the primarily ionospheric models of

Vuitton et al. (2006, 2007), or to the lack of ion-neutral chemistry in the case of earlier models. Lavvas et al. (2008 a, b); De La Haye et al. (2008); Lara et al. (1996); and Yung et al. (1984) all underestimate the abundance of the C3 group. Moreover the estimates of benzene mixing ratios vary sig-nifi cantly. Lavvas et al. (2008a,b) and Wilson and Atreya (2004) fi nd mixing ratios from a few times 10−10 to 1 × 10−9, while Krasnopolsky (2009) and De La Haye (2008) fi nd values in the 10−7 range. Vuitton et al. (2007) calculate a ratio of 3.0 × 10−6, which comes closest to observations, possibly because of their inclusion of a currently undefi ned 2-body reaction process.

Two recent papers have addressed the structure of the haze particles based on Cassini–Huygens observations. Liang et al. (2007) used UVIS ultraviolet observations of the continuum near 190 nm to infer the altitude distribution and size of the haze scattering the incident solar fl ux. Their analysis indicated that the upper atmosphere near 1,000 km contained up to 104 cm−3 macromolecules with sizes of the order of 10–20 nm. The photochemical calculations of Liang et al. suggest that polyyne polymers play a major role in the haze formation process. Lavvas et al. (2009) combined the UVIS ultraviolet observations with the ISS visible scattering observations to address the relationship between the detached haze layer and the haze layers extending to the surface. They found that the detached haze layer is formed as a result of changes in the particle sedimentation rate with

altitude. They assert that the observed mass fl ux 2.7 to 4.6 × 10−14 g cm−2 s−1, which is approximately what McKay et al. (2001), needed to explain the main haze layer, originates from production mechanisms dominated by upper atmospheric processes fi rst identifi ed by Waite et al. (2007). They therefore suggest that the main haze layer in Titan’s stratosphere is formed primarily by sedimentation and coagulation of par-ticles in the detached layer. More recently Sittler et al. (2009) suggest that suffi cient fullerenes are formed to con-tribute ~7% to the total aerosol infall.

Figure 8.10 summarizes the post Cassini–Huygens under-standing of the conversion of methane and nitrogen to organic macromolecules, then to organic aerosols, and even-tually to organic materials on the surface of Titan. The pro-cess begins with the dissociation and ionization of methane and molecular nitrogen by solar ultraviolet radiation and energetic particles in the upper atmosphere of Titan. This process sets in motion a rich ion neutral chemistry that pro-duces heavy positive ions and neutrals that eventually form macromolecules, many of which are negatively charged. Macromolecules precipitate rapidly from the formation layer near 1,000 to 550 km where they reach a size of 40 nm. They then slow through atmospheric viscosity to form the detached haze layer (Lavvas et al. 2009, 2008a,b). Some additional radical chemistry occurs during this descent, but most of the formation chemistry has already taken place in the low-pres-sure upper atmosphere as a result of ion neutral chemistry. Below 550 km the particles began to coagulate into the main aerosol layer and some additional chemistry likely takes place in this growth process. The cold temperatures of the troposphere lead to condensation of ethane and other organ-ics onto the aerosols (Hunten 2006), before they fi nally pre-cipitate onto the surface. The loss of hydrogen to space guarantees that the process will irreversibly convert methane in the atmosphere into organic residue on times scales from 10 to 70 million years (Mandt et al. 2009).

8.3 Conclusions: Laboratory Simulations and the Future of Titan Exploration

No laboratory simulations have been published that explic-itly try to match the Cassini data. However, it is instructive to look at published results from Titan simulations and com-pare them to the INMS spectra. For example, comparison of the results of Thompson et al. (1991) to the INMS results of Waite et al. (2007) shows that there is a broad similarity in the two results and suggests that properly conducted labora-tory simulations will be able to reproduce the INMS spec-trum (Fig. 8.11).

The laboratory mass spectrometer measurements of Imanaka et al. (2004) are compared to the mass peaks identi-


Tab

le 8

.4

Com

pari

son

of m

ixin

g ra

tios

from

the

INM

S w

ith s

ever

al m

odel

s

Spec

ies

Mag

ee e

t al.

(200

9)K

rasn

opol

sky

(200

9)C

arra

sco

et a

l. (2

008)

Lav

vas

(200

8a)

De

La

Hay

e et

al.

(200

8)V

uitto

n et

al.

(200

7)V

uitto

n et

al.

(200

6)W

ilson

and

A

trey

a (2

004)

Toub

lanc

et a

l. (1

995)

Lar

a et

al.

(199

6)Y

ung

et

al. (

1984

)

C2H

2M

ax3.

43 ×

10−

44.

2 ×

10−

43.

2 ×

10−

31.

4 ×

10−

42.

2 ×

10−

43.

0 ×

10−

45.

4 ×

10−

43.

6 ×

10−

31.

1 ×

10−

37.

4 ×

10−

3

Min

3.42

× 1

0−4

3.8

× 1

0−4

1.5

× 1

0−3

1.4

× 1

0−4

9.1

× 1

0−5

4.2

× 1

0−4

3.1

× 1

0−3

6.1

× 1

0−3

C2H

4M

ax3.

97 ×

10−

47.

0 ×

10−

42.

3 ×

10−

34.

5 ×

10−

46.

8 ×

10−

41.

0 ×

10−

36.

0 ×

10−

31.

5 ×

10−

32.

8 ×

10−

31.

5 ×

10−

35.

4 ×

10−

3

Min

3.91

× 1

0−4

6.0

× 1

0−4

1.1

× 1

0−3

4.4

× 1

0−4

9.1

× 1

0−5

1.3

× 1

0−3

2.8

× 1

0−3

4.2

× 1

0−3

C2H

6M

ax6.

05 ×

10−

51.

2 ×

10−

42.

2 ×

10−

49.

5 ×

10−

51.

0 ×

10−

42.

1 ×

10−

41.

1 ×

10−

53.

5 ×

10−

41.

1 ×

10−

3

Min

4.57

× 1

0−5

1.1

× 1

0−4

5.5

× 1

0−5

2.4

× 1

0−5

8.8

× 1

0−6

2.9

× 1

0−4

9.6

× 1

0−4

HC

NM

ax2.

44 ×

10−

46.

8 ×

10−

41.

5 ×

10−

39.

3 ×

10−

42.

1 ×

10−

42.

0 ×

10−

42.

0 ×

10−

47.

1 ×

10−

44.

6 ×

10−

42.

3 ×

10−

3

Min

2.40

× 1

0−4

5.3

× 1

0−4

7.1

× 1

0−4

7.5

× 1

0−4

5.8

× 1

0−5

5.0

× 1

0−4

4.7

× 1

0−4

1.5

× 1

0−3

C3H

6M

ax3.

45 ×

10−

61.

3 ×

10−

66.

8 ×

10−

6

Min

2.33

× 1

0−6

6.9

× 1

0−7

4.8

× 1

0−6

C3H

8M

ax4.

38 ×

10−

65.

2 ×

10−

64.

9 ×

10−

81.

5 ×

10−

77.

2 ×

10−

66.

7 ×

10−

77.

6 ×

10−

7

Min

2.87

× 1

0−6

2.4

× 1

0−6

1.2

× 1

0−8

6.5

× 1

0−8

3.1

× 1

0−6

3.4

× 1

0−7

4.1

× 1

0−7

C4H

2M

ax5.

65 ×

10−

62.

4 ×

10−

52.

5 ×

10−

61.

6 ×

10−

61.

0 ×

10−

56.

0 ×

10−

53.

7 ×

10−

61.

4 ×

10−

51.

6 ×

10−

51.

5 ×

10−

4

Min

5.55

× 1

0−6

1.1

× 1

0−5

8.8

× 1

0−7

7.3

× 1

0−7

3.1

× 1

0−6

6.4

× 1

0−6

7.9

× 1

0−6

7.8

× 1

0−5

C2N

2M

ax2.

20 ×

10−

62.

2 ×

10−

92.

4 ×

10−

82.

1 ×

10−

73.

4 ×

10−

5

Min

2.14

× 1

0−6

1.4

× 1

0−9

2.1

× 1

0−8

1.0

× 1

0−7

1.9

× 1

0−5

HC

3NM

ax1.

54 ×

10−

69.

8 ×

10−

62.

3 ×

10−

74.

0 ×

10−

52.

0 ×

10−

52.

3 ×

10−

61.

6 ×

10−

52.

3 ×

10−

4

Min

1.48

× 1

0−6

4.1

× 1

0−6

1.8

× 1

0−7

1.9

× 1

0−6

5.2

× 1

0−6

1.2

× 1

0−4

C2H

3CN

Max

4.39

× 1

0−7

8.0

× 1

0−6

1.0

× 1

0−5

1.0

× 1

0−5

1.1

× 1

0−6

Min

3.46

× 1

0−7

4.3

× 1

0−6

1.4

× 1

0−6

C2H

5CN

Max

2.87

× 1

0−7

3.6

× 1

0−7

5.0

× 1

0−7

5.0

× 1

0−7

Min

1.54

× 1

0−7

1.9

× 1

0−7

C6H

6M

ax2.

50 ×

10−

62.

7 ×

10−

74.

6 ×

10−

106.

1 ×

10−

73.

0 ×

10−

61.

4 ×

10−

9

Min

2.48

× 1

0−6

1.4

× 1

0−7

5.7

× 1

0−9

1.3

× 1

0−7

2.2

× 1

0−10

C7H

8M

ax5.

37 ×

10−

82.

7 ×

10−

72.

0 ×

10−

7

Min

2.51

× 1

0−8

1.3

× 1

0−7

Sour

ce: F

rom

Mag

ee e

t al.

(200

9), w

ith p

erm

issi

on f

rom

Els

evie

r.


Table 8.5. Comparison of CIRS data with several models.

CIRS data Models

Vinatier (2007)

Teanby (2007)

Krasnopolsky (2009)

Lavvas (2008)

Wilson and Atreya (2004) Lara (1996)

Toublanc (1995)

Yung (1984)

C2H

2Max 4.9 × 10−6 4.9 × 10−6 1.2 × 10−5 2.9 × 10−6 7.4 × 10−6 3.1 × 10−5 2.1 × 10−5 7.4 × 10−5

Min 3.6 × 10−6 1.5 × 10−5

C2H

4Max 2.0 × 10−7 1.3 × 10−7 3.3 × 10−8 4.6 × 10−8 1.5 × 10−8 2.3 × 10−8 1.5 × 10−7

Min 2.2 × 10−8

C2H

6Max 2.3 × 10−5 3.1 × 10−5 2.2 × 10−5 1.9 × 10−5 6.5 × 10−5 4.4 × 10−5 2.3 × 10−4

Min 2.2 × 10−5

HCN Max 1.3 × 10−6 6.2 × 10−6 3.7 × 10−6 2.0 × 10−6 1.3 × 10−6 2.2 × 10−6 9.7 × 10−6

Min 5.9 × 10−7 1.4 × 10−6

C3H

8Max 5.5 × 10−7 5.6 × 10−6 1.5 × 10−6 9.6 × 10−7 1.5 × 10−5 9.6 × 10−6

Min 6.2 × 10−7

C4H

2Max 4.5 × 10−9 1.6 × 10−7 1.0 × 10−8 2.9 × 10−8 1.2 × 10−6 8.0 × 10−8 6.1 × 10−8

Min 2.8 × 10−8

HC3N Max 5.4 × 10−7 3.4 × 10−6 8.7 × 10−9 2.2 × 10−8 2.2 × 10−8 9.8 × 10−7

Min 1.9 × 10−8

C6H

63.6 × 10−9 2.1 × 10−10 8.4 × 10−11 8.9 × 10−10

Fig. 8.10 Cartoon illustrating tholin formation as result of high-altitude ion-neutral chemistry. Over tens of millions of years methane is, with the loss of hydrogen to space, irreversibly converted to the complex hydrocar-bon-nitrile compounds that are the precursors of the aerosols (adapted from Waite et al. (2007) with permission from AAAS)

Fig. 8.11 Comparison of INMS spectrum and laboratory simulations. Solid line and red symbols are INMS observations of Titan’s atmosphere at an altitude of ∼1,000 km (Waite et al. 2005). Blue and gray symbols are laboratory results at 24 and 1,700 Pa respectively. The general comparability of the different results suggests that properly conducted laboratory simulations will be able to assist in understanding the INMS spectrum.


fi ed in the IBS ion mass spectrum in Fig. 8.12. While they are suggestive, the results indicate that the very low pressure ion-neutral chemistry present in Titan’s upper atmosphere has not been adequately explored in the laboratory and should be the focus of future investigations. Sittler et al. (2009) have carefully examined laboratory processes that lead to fullerence formation and their results suggest further lines of inquiry. The richness of Titan’s chemical environ-ment discovered by Cassini, and the complexity of chemical processes being explored in the laboratory together point to the need for further exploration in both arenas. In the case of Cassini this will take the form of an additional ∼60 encounters with the fascina ting atmosphere of Titan during the Cassini Solstice Mission.

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Chapter 8High-Altitude Production of Titan’s Aerosols8.1 Cassini Observations of Heavy Hydrocarbons in Titan’s Upper Atmosphere8.2 New Chemical Models Based on the Cassini Results8.3 Conclusions: Laboratory Simulations and the Future of Titan ExplorationReferences

chapter 8 high-altitude production of titan’s...

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